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Towards Understanding the Solar Atmosphere

Towards Understanding the Solar Atmosphere. by G. A. Doschek presented at the 6 th Solar-B Science Meeting 8-11 November 2005 Kyoto, Japan. Fundamental Questions 1- What drives solar activity? 2- What structures and heats the solar atmosphere?

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Towards Understanding the Solar Atmosphere

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  1. Towards Understanding the Solar Atmosphere by G. A. Doschek presented at the 6th Solar-B Science Meeting 8-11 November 2005 Kyoto, Japan

  2. Fundamental Questions 1- What drives solar activity? 2- What structures and heats the solar atmosphere? 3- How does solar energy affect the heliosphere and the Earth’s atmosphere? 1 x 105 K 5,800 K 2.0 x 10 K inner corona 2.0 x 106 K outer corona 2.0 x 106 K solar wind

  3. The Connectivity of Different regions of the Solar Atmosphere • The transition region is unresolved • Spicule-like structures are seen on the limb • The disk features vary approximately from loop-like structures to small spots • The actual emitting regions are only a few tenths of an arc-second in size • How does the lower transition region (T < 2.0 x 105 K) interface with the upper transition region and corona (2 x 105 K – 106 K+)? • Cool loops and coronal funnels? • Explanation of differential emission measure distribution • In this work I approach this question by attempting to correlate upper and lower transition region dynamics, specifically flows, by analyzing spectra from SUMER/SOHO. If there is connectivity between the upper and lower transition region, then in the simplest case radial flows in each region should be in the same direction. This is ongoing work, just recently begun. • C IV (1 x 105 K), S V (1.6 x 105 K), and Ne VIII (6.3 x 105 K) lines appear on a single SUMER wavelength window, assuring simultaneous and perfectly spatially aligned observations • Ne VIII images still show the network, allowing at least an approximate association with lower transition region features.

  4. The Magnetic Structure of the Solar Atmosphere

  5. A Model of the Transition Region/Corona Interface This schematic illustrates how the “standard model” can explain the appearance of solar structures at different temperatures

  6. Average Temperature-Density Structure of the Quiet Sun Chromosphere, Transition Region, and Corona The “standard model” of the solar atmosphere. Does it apply to all solar structures? The conductive flux makes the transition region extremely thin: Fc = aT5/2 dT/dh Ne VIII S V C IV

  7. Differential Emission Measure of the Quiet Sun Ne VIII (5.8) C IV

  8. Model of Sub-Resolution Filamentary Structure of Transition Region Spicules at 105 K Solar spectra indicate that the transition region contains fine structures with dimensions ranging from small values of about 1” (430 miles) down to very small values of about 0.043” (19 miles). Modeling of this fine-structure is now a hot topic in solar physics.

  9. How a Stigmatic Spectrometer Works NRL-NASA/HRTS rocket spectrum -C IV (~ 1549 Å). C IV is formed at 105 K in the solar atmosphere. EIS will obtain similar spectra with its narrow 1” and 2” slits. Images can be built up by rastering. Wavelength

  10. SUMER/SOHO Image and Spectra in the Light of O VI (3.2 x 105 K), N V (2.0 x 105 K), and O V (2.5 x 105 K) The O IV image was obtained at Sun center and covers an area of about 4 x 1010 km2. Each pixel has an area of 5 x 105 km2 and contains a measurement of both bulk and random plasma motion. This image contains about 80,000 measurements of dynamical parameters. Images in other spectral lines formed at different temperatures are available for comparison and correlation. CAT scan-like slice through the transition region in light of N V and O V.

  11. Optical Depth in Spectral Lines • Optical depth modifies line profiles by • resonance scattering • The shape of the line profile allows an • estimate of the true path length of emitting • plasma in the solar atmosphere along a • given line of sight. • The actual line shape depends on details • of the geometry. • Studies of line profiles in the lower transition • region indicate path lengths of only 50- • 100 km, or about 0.1 arcsec. Some details for lovers of radiative transfer: tau = F(Te) A Ne L f = fraction of total opacity in which the source function is zero Dashed profiles are optically thin lines: Gaussians O III line profiles

  12. O II 834.465 Å: opacity/km = 1.08 x 10-2 O III 833.749 Å: opacity/km = 8.22 x 10-4

  13. SUMER/SOHO 23 March 1996, about 19 hr UT 69” E – 268” W; 50” N – 350” N 297 spectra, 21 s exposures Slit: 1” x 300” Scan: East to West in steps of 1.125”. Center of slit: 200” North of equator.

  14. EIT - 171 EIT - 195 EIT - 304 EIT - 195

  15. Analysis Procedure • The three spectral lines are: C IV 1548.21 Å (1.0 x 105 K), S V 1572.98/2 Å (1.6 x 105 K), and Ne VIII 1540.85/2 Å (6.3 x 105 K). • Sum the Ne VIII, S V, and C IV line profiles for all the spectra in the three images, weighted by the total line intensities. There is one spectrum per line per 1”x1” pixel. Each image contains 93,852 spectra, of which 17,292 are statistically useful. • Determine by eye the central wavelengths of the three lines in the summed spectra. This is a somewhat arbitrary reference wavelength for defining redshifts and blueshifts. • Fit the three line profiles in each individual spectrum with Gaussian profiles, subject to the condition that the total intensities of all three lines exceed a certain value. This is done to ensure reasonable statistics. Unfitted spectra due to not meeting this condition are defined as having zero Doppler shifts. • The wavelengths of the lines determined by the Gaussian fits are subtracted from the reference wavelengths to yield the Doppler shifts.

  16. Summed weighted spectra for the raster images. The vertical lines are the central wavelengths estimated by eye.

  17. r

  18. dashed profile = Gaussian fit shift = 6.3 km/s

  19. S V C IV Doppler shifts in Å Ne VIII Downflow in Ne VIII (6.3 x 105 K) and S V (1.6 x 105 K). There is no evidence of a downflow in C IV (1.0 x 105 K). The downflow appears to be at the footpoint of what could be described as a plume in Ne VIII. A similar downflow appears in another Ne VIII plume-like feature. The down- flow is quite small, about 6.3 km/s.

  20. Doppler Shifts in Angstroms C IV S V C IV The C IV spectrum is one of the spectra that pass through the upflow region, showing that it is a real feature, often seen in HRTS spectra. Ne VIII C IV 0.1 Å≈20 km/s at 1548 Å 0.17 A ≈ 32 km/s

  21. Hist.: C IV Solid: S V Dashed: Ne VIII 0.17 Å ≈ 33 km/s

  22. Doppler shifts in Angstroms C IV S V The entire image: only data with total line intensities greater than 100 counts are considered. Ne VIII + 0.05 Å is a result of changing the reference Ne VIII wavelength by +0.05 Å. Ne VIII Ne VIII + 0.05 Å

  23. S V – C IV S V – Ne VIII C IV – Ne VIII Doppler shift differences: Note that perfect agreement between Doppler shifts gives zero net differences. Only S V – C IV approximates this result.

  24. Conclusions: The Lower Transition Region • The lower transition region is unresolved (at the best orbiting spacecraft data now available (1-2 arcsec imaging). • Different temperature regions are correlated in intensity with correlation coefficients of ~ 0.9 • Non-thermal mass motions correlation coefficients are ~ 0.5-0.8 • Path lengths from optical depth estimates are 50-100 km • High spatial spectrometers like VAULT are needed to make further progress in understanding the smallest structures in the transition region. SUMER images show many elongated structures, suggesting that cool loop models are indeed at least partly applicable to the real Sun. • Future Directions: Continue to attempt to use SUMER spectra to determine the percentage of cool loops relative to “classical” transitions regions in the solar atmosphere; analyze results from continuing VAULT rocket program; use He II and Si VII lines in EIS Solar-B to attempt a similar analysis.

  25. Solar Flare Science Issues • Solar flare flux tubes at multimillion degree temperatures do not have the appearance expected from 1D numerical simulations. Why? X-ray spectral line profiles do not show evidence for expected “chromospheric evaporation”. Why not? What are the implications for the standard reconnection models of flares? • X-ray images of solar flares at multi-million degree temperatures can show surprising morphological complexity that is difficult to interpret. Not all flares look like simple loop arcades (“Masuda”; “Bastille Day” events) or long-duration events (“Tsuneta flare”). • What are the physical conditions in the reconnection region above the soft X-ray flare loops? • Can EIS “see” hard X-ray emission as well as the soft X-ray signatures?

  26. Ca XIX Resonance Line X-ray Spectra

  27. A Multi-Thread Flare Loop Model • The overall soft X-ray magnetic envelop is assumed to be composed of sub-resolution magnetic “threads”. • The threads are modeled as individual flare loops using the NRL 1D solar flux tube model. • The flare onset is modeled as a succession of independently heated threads. • The length of the flare loop is determined by observations of the overall magnetic envelop. • The energy deposited in each thread and its cross-sectional area are related to the GOES fluxes (Warren & Antiochos 1994). The energy in each successively heated thread is determined such that the X-ray flux matches the GOES light curves. • The BCS spectrometers on Yohkoh serve as a completely independent test of the model. • The BCS data support the multi-thread model (Warren & Doschek 2005, ApJ, 618, L157).

  28. Multi-Thread Results for Yohkoh/BCS • Observe evolution and distribution of multi-temperature plasma • Requires images that span 104 – 107 K. From Warren & Doschek, ApJ, 2005

  29. EIS: The Use of Slits and Slots Solar flare time sequence Skylab flare spectral images and EIS 40” wide slot Slit spectra give line profiles

  30. Images of Multi-million Degree Flare Loops • Do flare loops at temperatures of 12-25 MK look like what we expect??? • No, they don’t, but as they cool to 1-3 MK they look more and more like respectable 1D loops should. TRACE observations confirm the gross flare loop morphology seen by Yohkoh.

  31. Solar Flare Reconnection Model Solar flare 45,000 km This schematic flare model provides theoretical guidance for analyzing solar flare data.

  32. Observe motions of ejecta and inflows above arcade Requires cooler lines and longer exposures – summing of images Other Objectives: Reconnection Region From Shibata et al., ApJ, 1995 From Sheeley, Warren, and Wang, ApJ, 2004

  33. Conclusions • The first observed BCS Ca spectra are not blue-shifted in disk flares. However, the multi-thermal thread model can reasonably reproduce the observations. A “sudden impact” evaporation model cannot reproduce the observations. • TRACE images indicate the presence of substructure (i. e., threads in soft X-ray flare loops (see Doschek, Strong, & Tsuneta, 1995, ApJ, 440, 370 for a discussion of Yohkoh flare loops). • BCS Ca XIX/S XV initially observed electron densities are about 3.5 x 1010 cm-3 and emission measures are about 6 x 1046 cm-3, assuming the multi-thread scenario and a nominal electron temperature of 12 x 106 K (see Doschek & Warren, ApJ, 629, 1150). • The appearance of multi-million degree flare loops exhibits hot knots at the loop-tops and sometimes asymmetric loop brightening. The knots and asymmetric brightening are not projection effects and need to be explained in terms of our conventional views of plasma confinement in magnetic flux tubes. • The soft X-ray flare onset for some events begins before the hard X-ray flare onset. • EIS can use He II 256 Å emission as a hard X-ray proxy for footpoint emission.

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